Think back to the lab on Atomic Spectra and you will remember the discussion of spectral classes that is repeated below. Now it is time to put the tools of spectroscopy together with what we know about the nature of stars.

Spectral Classes for Main Sequence Stars

In 1872, Henry Draper first photographed stellar spectra. This represented a tremendous advance. Instead of sketching spectra astronomers could directly record, compare and measure them. Spectra of thousands of stars became available for precise analysis. The first classification scheme was based on the intensity of a star's hydrogen emission lines. This system was arranged in alphabetical order. 'A' type stars had the strongest hydrogen lines and 'O' the weakest. After physicists came to understand the nature of atoms in the 1920s it became necessary to "rearrange" the alphabet.

This new system was begun by Edward Pickering and completed by Annie J. Cannon and a group of young women assistants who invented a system of spectral classes based on the number and appearance of spectral lines. When Annie Cannon published the Henry Draper catalogue it contained spectral data on 225,320 stars and became the basis for all modern astronomical spectroscopy.

The sequence of spectral classes now begins with the hottest stars, the O stars, and ends with the coolest stars, the M stars. This is summarized in the following table. The strong lines column describes the prominent line features in the spectrum. In the next lab you will study the properties of stars. Remember that a star's classification (O, B, A, F, G, K, or M and the 10 subdivisions within) depends on its color and the pattern of absorption lines, both of which are determined by its temperature. For the subdivisions 0 - 9, 0 is the hot end and 9 is the cool end.
 
 
 
 


 
 
Principal Classes And Characteristics
  Class   Color       Temperature  Strong Lines 
O blue-white 35,000 K ionized helium
B blue-white 21,000 K  neutral helium
A white 10,000 K hydrogen
F green-white  7,000 K ionized calcium
G yellow 6,000 K  ionized calcium
K orange 4,500 K neutral sodium
M red 3,000 K titanium oxide

 
 
    Before, we begin, note that you can represent spectra in two different ways.  One is how we did it in the Atomic Spectra Lab (on the left below).  The second, which astronomers usually use, is to make a graph of brightness vs wavelength (on the right below).  The spectra we will look at below are shown in this second way.

Now look at these reference spectra for several different spectral types.   Instead of showing them as color spectra, we make a plot of brightness vs. wavelength.  So for instance, the hot O star is brighter at the blue end, and the cool M star at the red end.  The dark absorption lines show as dips in the plot.  Keep that window up as you will need to compare spectra to these in order to determine the spectral types of several stars.

Let's look at the distinguishing characteristics of each spectral type, as illustrated by the reference spectra.

The figure shows spectra for seven stars of decreasing temperature: O5, B3, A6, F6, G7, K5 and M4. Note the following characteristics:

1. As temperature decreases, the spectrum shifts from peaking towards the short wavelength (blue) end, to the long wavelength (red) end. This reflects the overall, rough blackbody shape of stellar spectra, and Wien’s Law for blackbodies (see Spectroscopy section of Atomic Spectra lab).

2. Hotter stars have more energetic photons that cooler stars. One consequence of this is for helium. It takes rather energetic photons to raise helium atoms to excited levels and thus produce an absorption line. Thus, helium absorption lines are only present in the hottest stars, even though helium itself is present in all stars. By the same reasoning, the coolest stars do not show absorption lines of hydrogen, even though they are mostly hydrogen! They do not have many photons energetic enough to raise hydrogen atoms to excited levels.

3. But hotter stars don’t show hydrogen lines either. Why? The reason is ionization. Energetic photons in hotter stars also cause more ionization. A consequence is that hydrogen is ionized in the hotter stars. Since it has lost its only electron, there are no electrons left to make transitions between energy levels and produce spectral lines. Calcium is easy to ionize, and even moderate temperature stars have energetic enough photons to do it. The pair of calcium lines is from ionized calcium atoms that have lost one electron, and are seen in moderate temperature stars. They get weaker in the hottest stars because such stars have energetic enough photons to cause calcium to lose more electrons. The sodium line is from neutral sodium and is seen in cooler stars, because even they can cause neutral sodium atoms to be in excited states. In moderate and hot stars, sodium is ionized, so we don’t see the neutral sodium line. The same is true for the neutral magnesium line marked.

4. In the coolest stars, atoms move relatively slowly and as a consequence can bond together when they collide, producing molecules. Molecules have rich absorption patterns of their own, often producing very broad features. Prominent, broad trough-like features due to the TiO molecule are seen in the coolest stars.

Now for each of the five stellar spectra below, decide which of the reference spectra in the figure it most closely resembles. Consider the overall shape and the spectral lines present.  If it appears to be between two types, note this too and note which one it is most like. From the table of spectral types and temperatures, roughly estimate the temperature of each star.